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Shock-induced PDR in the Herbig-Haro object HH 2
B. Lefloch, J. Cernicharo, S. Cabrit, D. Cesarsky
To cite this version:
B. Lefloch, J. Cernicharo, S. Cabrit, D. Cesarsky. Shock-induced PDR in the Herbig-Haro object
HH 2. Astronomy and Astrophysics - A&A, EDP Sciences, 2005, 433, pp.217-227.
�10.1051/0004-6361:20042093�. �hal-00398146�
1 Laboratoire d’Astrophysique de l’Observatoire de Grenoble, BP 53, 38041 Grenoble Cedex, France
e-mail: lefloch@obs.ujf-grenoble.fr
2 Consejo Superior de Investigaciones Científicas, Instituto de Estructura de la Materia, Serrano 123, 28006 Madrid, Spain 3 LERMA, Observatoire de Paris, UMR 8112, France
4 Max-Planck Institut für Extraterrestrische Physik, 85741 Garching, Germany
Received 25 May 2004/ Accepted 24 November 2004
Abstract.We report mid-infrared (5−17 µm) and SO, CO,13CO millimeter line observations of the protostellar jet HH 2 and
the parental molecular cloud. We have detected for the first time mid-infrared emission along a protostellar jet. We find that the outflowing gas extends much further away than the Herbig-Haro object HH 2, showing direct evidence that downstream gas has been accelerated by previous outflow events. These gas layers appear to have been detached from the parental cloud, as they are distributed around a cavity, probably dug by protostellar outflow(s). SO emission is detected in shocked gas regions associated with outflows. The UV field produced in the strong shock region HH 2H-A has produced a low-excitation Photon-Dominated Region at the walls of the cavity, which is detected in the PAH emission bands and in the continuum between 5 and 17µm. This continuum arises from very small grains transiently heated by a FUV field G 20 G0, which probably formed from evaporation
of dust grain mantles in shocks.
Key words.ISM: Herbig-Haro objects – ISM: individual objects: HH 1/2 – ISM: jets and outflows – ISM: dust, extinction – stars: formation
1. Introduction
The Herbig-Haro system HH 1–2 is one of the best-studied star forming region. It attracted early attention because of the spec-tacular shock emission regions HH 1 and HH 2 that trace the in-teraction of the protostellar jet with the ambient cloud. This jet is characterized by rather high velocities in the ionic and atomic material (up to 480 km s−1) and large proper motions in the im-pact region, as measured with HST (Bally et al. 2002). The optical jet is associated with a molecular outflow, first mapped in CO by Moro-Martín et al. (1999). The jet makes a strong in-clination angle with the line of sight≈80o. The high-velocity,
collimated component of the outflow (the “molecular jet”) cov-ers deprojected velocities of 15−80 km s−1. This system was first studied by Martin-Pintado & Cernicharo (1987) who sug-gested that the spatial distribution of the molecular emission resulted from the interaction of the outflowing jet with the am-bient molecular cloud, producing cavities and density enhance-ments along the walls of these cavities.
In the HH 2 region, the impact in the ambient gas results in shocks with a wide range of velocities: from 20−30 km s−1 in the rim, corresponding to the individual knots E-K (Lefloch et al. 2003) up to 180 km s−1in the Mach disk region, which was identified as the optical knots HH 2H-2A. Observations in the UV (Boehm-Vitense et al. 1982; Raymond et al. 1997) and X-ray (Pravdo et al. 2001) show extended emission from this shock region.
There has been a long debate as to whether such a high-energy field would have any noticeable impact on the local sur-roundings. Based on previous molecular line observations, the gas downstream of HH 2 appears to be in a quiescent state. However, Davis et al. (1990) report an enhanced abundance of HCO+in that region. Torelles (1992) obtained a similar re-sult for NH3. A detailed study by Girart et al. (2002) revealed
the peculiar chemical composition of the molecular emission peak in the region downstream of HH 2: some molecules were found largely overabundant (CH3OH, H2CO, HCO+, SO, SO2)
and others, like CS and HCN, were underabundant. These fea-tures appeared to agree qualitatively with chemical models of UV-irradiated gas clumps (Wolfire & Königl 1993; Viti & Williams 1999), which predict large abundance enhancements in a wide variety of molecular species, from either purely gas phase reactions or the release from icy grain mantles. Time-dependent modelling by Viti et al. (2003) could account for
some of the molecular abundances measured by Girart et al.
(2002), paying special attention to the HCO+emission. Other observational evidence of the impact of the UV field on the am-bient gas around HH 2 came from the observation of the FIR lines [CII] 158µm and [OI] at 63 and 145µm with ISO/LWS at 80resolution by Molinari & Noriega-Crespo (2002). They found that shock models cannot account for the line intensities measured and concluded that at least part of the emission must arise from a PDR associated with HH 2, but lack of angular
resolution prevented any further characterization. Their mod-elling also made the assumption that the PDR is excited mainly by 2γ UV photons.
On the other hand, it is well established that shocks can strongly alter the chemical composition in the entrained gas of protostellar outflows (see e.g. Bachiller & Perez-Gutierrez 1997). Indeed, recent HCO+J= 1 → 0 interferometric
obser-vations by Dent et al. (2003) reveal a high-velocity component moving along HH 2. The authors conclude that the observed HCO+abundance enhancement is consistent with shock chem-istry in the turbulent mixing layer associated with the jet.
In this article, we reassess the nature of the molecular emission around HH 2, from 5−17 µm spectro-imaging ob-servations of the HH 2 region obtained with the ISOCAM camera onboard ISO and complementary observations of mil-limeter transitions of CO,13CO and SO, at the IRAM 30 m tele-scope, paying special attention to the position studied by Girart et al. (2002), which we will refer to as the “molecular emission peak”. Our observations show HH 2 to be even more complex than initially thought. We find that previous outflow episodes have accelerated gas ahead of HH2 and have dug a cavity in the parental cloud. Shocked gas is detected along the walls of the cavity. The mapping of SO emission shows that it is as-sociated with the shock interaction of the outflowing gas with the parental cloud. The strong J-type shocks associated with HH 2H-A have induced the formation of a Photon-Dominated Region (PDR) at the inner wall of the cavity, which could be mapped and characterized in the mid-IR.
2. Observations
The CO and13CO millimeter line data were obtained at the IRAM 30m telescope and have been presented and discussed in Moro-Martín et al. (1999), hereafter MM99.
The SO transitions 34−23 at 138.17864 GHz and 23−12at
99.299883 GHz were observed with the IRAM 30m telescope in March 1993. The observing conditions were very good, with typical system temperatures of 180–200 K at 2mm. The angu-lar resolution of the telescope is 17and 24, respectively, at these frequencies. An autocorrelator with a spectral resolution of 20kHz was used as a spectrometer. The data was smoothed to obtain a kinematical resolution of≈0.05 km s−1. The SO emission was mapped with a 15 sampling over a region of about 200by 200centered on the VLA 1 source. The flux is expressed in units of main-bream brightness temperature. The efficiency of the telescope was 0.65 and 0.55 at 99.299883 GHz and 138.17864 GHz respectively.
The mid-infrared observations were obtained with the ISO satellite (Kessler et al. 1996) and the ISOCAM instrument (Cesarsky et al. 1996). The low resolution spectra (λ/∆λ = 40) between 5 and 17µm were obtained in revolution 691 with the Circular Variable Filter (CVF) with a pixel scale of 6and a to-tal field of view of 3centered on the HH 2 object. We present the data reduced with the pipeline version OLP10. The size (HPFW) of the Point Spread Function (PSF) is≈6for a pixel scale of 6.
Accurate astrometry (better than 2) was established us-ing a second CVF map containus-ing the optically visible
Cohen-Schwartz (CS) star, taken in revolution 873. Details are given in Cernicharo (2000) and Lefloch et al. (2003). The zo-diacal light and a possible large-scale emission across the field were suppressed by subtracting a reference spectrum from the whole dataset. A spectrum of the residual emission, away from HH 2, is shown in panel A of Fig. 8c (the “empty field”). In addition to the CVF map we obtained an image of the flux integrated in the range 5.0−8.5 µm, including the H2
pure rotational lines S(8)–S(4) and the PAH bands at 6.2 and 7.7 µm. We extracted a map of the continuum emission in the range 13.9−15.5 µm, outside the interval of emission of the [Ne 2], [Ne 3] and H2 S(2) lines. The data are presented in
Fig. 8. Coordinates are offsets (arcsec) relative to the position of VLA1:α2000= 05h36m22.6s,δ2000= −06◦4625.
The SNR of the data is not very high; we have averaged the signal over four fields typical of the region, which are drawn in Fig. 8: in the cloud (A), along the jet (D), and over the “ring” (B,C, see below). The interstellar extinction towards HH 2 was estimated by Hartmann & Raymond (1984), who measured typical reddenings E(B− V) = 0.11−0.44. Based on the ex-tinction curve of Rieke & Lebofsky (1985), it appears that the flux dereddening corrections are negligible and we use uncor-rected flux values in what follows.
3. Entrained gas in the HH 2 region
Although the region has been extensively studied and several outflows identified (see e.g. Chernin & Masson 1995; Correia et al. 1997) it is only recently (1999) that the molecular coun-terpart to the HH 1–2 jet was discovered (Moro-Martin et al. 1999). The confusion in the HH 2 region is indeed very high; because of the weakness of the outflow emission and a propa-gation close to the plane of the sky (about 10◦, Noriega-Crespo et al. 1991), it is difficult to identify the different kinematical components in line spectra.
Several observations suggest that the medium in front of HH 2 has been accelerated in the past. Henney et al. (1994) showed that the medium in front of the counterjet of HH 2 (HH 1) is moving at a considerable velocity, 200 km s−1. Ogura (1995) reported the presence of two giant bowshocks symet-rically located at 30 arcmin from the protostellar core and aligned with HH 1–2, ejected some 104yr ago. These
bow-shocks were interpreted as the signature of previous ejections from the source powering the HH 1–2 jet. We have re-analyzed more thoroughly the12CO data presented by MM99 to search for any hint of outflow/ejections older than the HH 1–2 jet mapped by them.
3.1. A CO outflow ahead of HH 2
We show in Fig. 1 the CO J= 2 → 1 high-velocity gas emis-sion in the HH 2 region, as mapped by MM99. The
red-shifted (vlsr= 12−16 km s−1) outflow wing associated with the
HH 2 jet extends between the driving source VLA 1 and the HH object (the H2 knots of shocked gas are marked by white
stars in Fig. 1. Close inspection of the data reveals another CO component, well collimated, which propagates downstream of the red wing at blueshifted velocities (vlsr = 2−4 km s−1).
Fig. 1. (Left) Map of the high-velocity outflow emission in the 12CO(2–1) line (thick black contours) superposed on the map of the 13CO(2–1) emission (thick white contours) and a [SII] image of the region (greyscale; Reipurth et al. 1983). The image has been saturated
to outline the ring and the weak emission downstream of HH2. The redshifted wing of the HH 2 outflow is drawn in dashed contours. Contours range from 4 to 20 K km s−1by step of 2 K km s−1. The blueshifted wing is drawn in solid contours. First contour and contour interval are 1 and 0.5 K km s−1respectively. The white stars mark the position of the HH 2 knots A to L. The position of the regions discussed in the text (the “ring” and the molecular emission peak downstream HH 2) are indicated. (Right) Map of the13CO(2–1) emission integrated between 5.5
and 6.75 km s−1(thick white contours) and between 9.5 and 12 km s−1(dashed) superposed on a [SII] image of the region (greyscale). Intensity contours range from 6 to 11 K km s−1.
A spectrum of the CO outflowing gas near the brightness peak at offset position (60,−75) is displayed in Fig. 2. The blueshifted wing overlaps very well the ambient gas layers, as traced by13CO. The maximum of brightness in the blue
com-ponent is detected in the region that coincides with the HCO+ high-velocity wing reported by Dent et al. (2003). This kine-matical component extends from knots HH 2 K-E and propa-gates over≈50, ahead of knot L (see Fig. 1). The optical [SII] emission of HH 2 reveals a faint jet that propagates ahead of HH 2 and coincides spatially with the CO blueshifted outflow (see left panel in Fig. 1). This jet is likely to be the driving source of the blueshifted molecular material.
A crude estimate of the parameters in the outflowing gas was made by assuming a kinetic temperature of 30 K. As can be seen in Fig. 2, the CO spectra indicate antenna temperatures of20 K, which is a lower limit to the intrinsic line brightness. The brightness temperature of the very optically thick CO lines provides a lower limit to the actual gas kinetic temperature, which is about 25 K. An upper limit is provided by the main-beam brightness temperature40 K; we adopt an intermedi-ate value of 30 K in this work. This value is characteristic of the ambient gas and may underestimate the actual temperature in the outflowing gas. It is higher than the 13 K estimated by Girart et al. (2002). These authors obtained this estimate in-directly, based on a multi-transition Monte-Carlo analysis of HCO+. We favor our approach, which relies directly on the observation of a “standard” tracer; excitation problems can be very severe for HCO+. The presence of a thermal gradient in the layers cannot be excluded, in which case the temperature
Fig. 2. Montage of12CO(2–1) and13CO(2–1) spectra at 3 positions:
the protostellar core (top), the “ring” (middle), and the “molecular peak” downstream of HH 2 (bottom).
at the surface (traced by the low-excitation CO line) could be higher than in the inner denser regions probed by HCO+.
Integrating the wing emission in the range 2−4 km s−1, we find a mass of 0.015 Mand a momentum of 0.1 Mkm s−1. These parameters compare well with those derived for the
Fig. 3. Velocity-integrated intensity channel map of the SO 34−23emission in the HH 1–2 region (contours), superposed on the 5.0−8.5 µm
emission map obtained with ISO. First contour and contour interval is 0.2 K. Velocity is indicated in the top left corner of each panel. The position of VLA 1 is marked by a white star. The individual HH knots are drawn in black. Coordinates are in arcsec offset with respect to VLA 1.
high-velocity molecular outflow associated with the HH 2 jet. This blueshifted component appears relatively well collimated and is aligned with the redshifted Herbig-Haro jet HH 2, which suggests a common origin for both ejections. Both components slightly overlap at the position of the knots E-F-K, where the maximum of shocked H2 emission is detected in the mid-IR
(Lefloch et al. 2003). The alignment and overlap is better seen in the higher-angular resolution observations of Dent et al. (2003). Though appealing, it is not clear if this change of ori-entation results from precession of the jet driving source (blue and red components would be tracing two different ejections) or from deflection upon impact on dense obstacle (the com-ponents are tracing one single jet). Both comcom-ponents have ap-proximately the same (projected) size≈50 = 3.3 × 1017cm, and could trace two ejections separated by 500 yr, assuming a typical ejection velocity of 200 km s−1. This blueshifted out-flow and the optical jet downstream of HH 2 provide direct ev-idence that material ahead of HH2 has already been accelerated by previous outflowing ejections.
3.2. The high-density gas emission
SO is a tracer whose emission is more highly contrasted than CO because its abundance can be greatly enhanced in shocks and in the high-temperature central protostellar regions (the “hot cores” of low- and high-mass protostars), where grain mantles are evaporated and S-bearing molecules released in the gas phase. SO is therefore especially suited to map the shocked gas of the outflow and its interaction with the ambient cloud. We detected extended SO emission over the entire sur-veyed region. We show in Fig. 3 a channel map of the velocity-integrated flux of the SO 34−23line. We found two maxima; the
first one is located at position (−75, 12), the second one lies
very close to knot HH 2L at position (60,−75), in very good agreement with the “molecular peak” position determined by Girart et al. (2002). At this position, SO lines are bright (∼4 K), more than twice as bright as in the protostellar core, where the HH 1 jet is detected as a blue wing, and in the HH 2 outflow (see Fig. 4).
We concentrate on the region downstream HH 2 in what follows. From the half-power contour, we estimate a transverse (deconvolved) size of≈20, hence barely resolved by our ob-servations. We show in Fig. 4 the profiles obtained in a 4-point cut in declination across the brightness peak at (60,−75). The line profiles peak at about 6.5 km s−1. However, they are characterized by several kinematical components, which vary between adjacent positions, separated by 15. Atδ = −90, a secondary component is detected at 6.89 km s−1. The veloc-ity channel map suggests that the 7 km s−1 component peaks at lower declination, near position (50,−90). Unfortunately, the data sampling is lower in that region (30) and does not al-low us to resolve details in the structure of this feature. North of the peak (δ > −60), a secondary component is detected at 10 km s−1. It is is rather weak, with Tmb ∼ 1 K, but it is
unam-biguously detected across the whole map. It is spatially asso-ciated with low-redshifted gas which extends from the VLA 1 protostellar core to HH 2 and beyond. Overall, the gas traced by SO does not appear quiescent in the region downstream of HH 2.
3.2.1. Physical properties of the SO emission
We have carried out an LVG analysis of the SO emission at a few positions close to the brightness peak region, consider-ing one sconsider-ingle gas component. We have neglected any thermal and/or density gradients in the emitting gas layer, despite the
Fig. 4. Spectra of the SO 34−23and 23−12lines observed towards the
molecular emission peak downstream HH 2 at offset position (60,
−75) and the protostellar source VLA 1. Flux is in main-beam
bright-ness temperature. Line intensities have been multiplied by 2.0 for the VLA 1 spectra.
fact that the SO emission arises from a shocked gas layer, as we discuss below in Sect. 3.2.2. Taking into account such an effect would introduce an additional degree of freedom in the modelling, for which the present observational data set does not bring any constraints. We adopt a kinetic temperature of 30 K, similar to that determined in the lower-density gas traced by CO (see above). Collisional rates for SO have been derived from those of CS-H2 computed by Green & Chapman (1978);
the procedure for3Σ molecules is described by Fuente et al.
(1990). We obtained very similar results using the more recent collisional rates computed by Green (1994) and extrapolated at temperatures lower than 50 K.
At the peak, both lines are bright with main beam tem-peratures of 4.2 K and 4.7 K for the 23−12 and 34−23
tran-sitions respectively. Linewidths are≈1 km s−1. We find a so-lution for a column density N(SO) = 5.6 × 1013cm−2 and
a gas density n(H2) = 5.5 × 105cm−3. The lines are
op-tically thin with opacities of 0.45 and 0.23 for the 34−23
and 23−12 respectively. The LVG calculations predict
in-tensities of 2.2 K for the 56−45 219.9494 GHz and 0.91 K
for the 67−56 261.8437 GHz, both transitions observed by
Fig. 5. Variations of the SO 34−23/23−12 (thick solid) and
67−56/56−45(thick dashed) intensity line ratios with temperature and
density at the molecular peak. The SO column density is taken to be 5.6 × 1013cm−2 (see text). Thin contours draw the values at 90
and 110.
Girart et al. (2002) at the SO peak. This is in good agreement with the line brightness, once corrected for the beam dilution factor (respectively 0.257 and 0.322).
We note that a lower temperature, of the order of 15 K for the emitting gas, would require much higher densities, of the order of n(H2)= 5 × 106cm−3, and N(SO) 6.5 × 1013cm−2.
These densities are far too high and lead to a flux much larger than what is observed for the 56−45 and 67−56 transitions
(2.7 K and 1.3 K respectively). These results differ somewhat from the simple LTE analysis by Girart et al. (2002). Indeed, the excitation temperature ranges from 24 K for the 23−12
tran-sition down to 9 K for the 67−56transition, showing that these
transitions are far from being thermalized.
Taking into account the uncertainties in the brightness de-termination, it appears that various sets of physical conditions can account for the observed lines. We have explored the range of solutions in the parameter space defined by the density and the temperature, adopting the SO column density determined above at the brightness peak (the “best solution”). We show in Fig. 5 the intensity line ratios 34−23/23−12 (thick solid line)
and 67−56/56−45 (thick dashed line). A good agreement is
obtained for densities in the range 3−6 × 105cm−3 and
tem-peratures between 25 and 40 K. Allowing a variation of 10% for each ratio, we find that the density has to be more than 105cm−3, and the temperature less than 100 K. If the
den-sity n(H2) is larger than 2.5 × 105cm−3 i.e. the average gas
density in the layer (see Sect. 4.1), the temperature is relatively well constrained, between 20 K and 60 K. Observations of higher-excitation transitions of SO (and CO) would help con-strain the temperature.
Comparison of the relative intensities at the neighbouring positions (Fig. 4) shows that the 34−23intensities is lower than
23−12, suggesting a change of excitation conditions, namely
a lower density. An LVG analysis yields a column density
N(SO) = 3.0−4.0 × 1013cm−2 and a density n(H
2) 2 ×
105cm−3. If the gas temperature were lower, typically 15 K,
the densities derived would be n(H2) 6−7 × 105cm−3,
which is very unlikely for structures with a size of 1017cm in
the molecular cloud. The densities derived from our LVG ap-proach are very similar to those quoted by Dent et al. (2003),
Fig. 6. Comparison of the emission distribution of the SO 34−23(black
contours) at various velocity intervals (6.5, 9.5, 10, 10.5 km s−1) with the CO 2–1 high-velocity outflow. First contour and contour interval are 0.2 K.
from millimeter dust continuum, and from modelling of the HCO+emission by Girart et al. (2002).
The secondary component at 10 km s−1 detected North of the SO peak (δ = −60andδ = −45) is characterized by high-excitation conditions with a ratio 34−23/23−12 ∼ 1.5. Such a
high value is found only at the SO emission peak ahead of HH 2 (it is less than 1 elsewhere in that area). Carrying out an LVG analysis of the position (60,−45) and adopting a kinetic tem-perature of 30 K, we find a gas density n(H2)= 5 × 106cm−3
and N(SO)= 2.0 × 1013cm−2. Again, the density estimate
re-lies on the temperature adopted, but we stress that 30 K is prob-ably a reasonable lower limit (lower temperatures require even higher densities,∼107cm−3, irrealistically large for the
large-scale entrained gas of a molecular outflow). The emitting gas is therefore characterized by high densities, much larger than in the ambient parental cloud, and possibly warmer temperatures. Taking into account that this gas is moving into the ambient cloud, we conclude that at least part of the SO component is tracing gas shocked and/or accelerated by the HH 2 jet.
3.2.2. Origin of SO emission
How does the distribution of the SO-emitting gas relate to the CO outflowing gas detected in HH 2? Figure 6 shows that SO traces the large-scale features associated with the accelerated gas detected in CO on each side of the protostellar core at 8.5 km s−1. In the protostellar core itself, the SO line profiles exhibit blue- and red-shifted wings, which are the signature of the HH 1–2 outflow. We did not find any evidence of emission
along the jet down to HH 2; instead, we detect in the same
ve-locity range emission that is spatially shifted with respect to the jet; such a shift could be an indication that the emission arises from the jet rim, i.e. the low-velocity shock region of the jet
(see Lefloch et al. 2003). The gas column densities derived are indeed comparable to the values encountered in other young protostellar outflows, like L1157 (Bachiller & Perez-Gutierrez 1997). Observations at higher-angular resolution would help clarify this point.
As discussed above, the high-velocity blushifted outflow propagates into the gas layers downstream of HH 2, where SO is detected at velocities equal or very close to ambient (Fig. 6). The SO emission is elongated perpendicular to the CO out-flow direction, which propagates across the SO region -limited by the contour at half-power- and finishes≈40Southeast of the SO peak. Both ambient SO and high-velocity CO bright-ness peaks are very close to each other (less than 10) and to knot HH 2L, where shocked molecular gas is detected (see e.g. Lefloch et al. 2003). This peak coincides with the maximum of density and of SO column density. Analysis of the physi-cal conditions shows that the gas is denser than in the ambient parental cloud and could be warmer than the 30 K estimated from CO. The detection of high-velocity gas at the position where the density is highest suggests that SO is tracing the im-pact of the flow in the cloud. Also, the amount of SO material detected along the HH 2 jet and the layers downstream of HH 2 are similar (to less than a factor of 3), which suggests a com-mon mechanism, and the column densities derived are typical of protostellar outflows.
Because the inclination of the flow with respect to the line of sight is very large, it is difficult to determine at which veloc-ity SO emission occurs when the flow impacts the gas layers. Detailed modelling would help constrain the physical condi-tions of the shock that could account for the SO emission de-tected or if on the contrary UV radiation plays an important role in the emission detected.
4. A cavity in the cloud
In the previous section, we showed evidence that the proto-stellar outflow from VLA 1 has propagated beyond HH 2. Its impact on the ambient gas causes shocks which are detected in high-density tracers such as SO. Previous studies of the gas and dust structure in multiple star forming regions suggest that flows can have some other impact on the cloud dynamics; the study of NGC 1333 revealed for instance the presence of cav-ities in the cloud, apparently dug by the protostellar outflows (see e.g. Lefloch et al. 1998). We study in this section the large-scale distribution of the molecular gas kinematics around HH 2, to search for any similar impact of the outflowing gas on the parental cloud.
4.1. The low-density gas around HH 2
We have studied the distribution of the low-density molecular gas by mapping the emission of the13CO(2–1) line. We present
in Fig. 7 the map of the velocity-integrated flux in the range 5−11 km s−1. Three regions can be distinguished depending on the strength of the emission: the ambient cloud, between 8 and 9 km s−1; a fragment downstream of HH 2, between 5 and 8 km s−1; the “filament” southwest of the protostellar core, be-tween 10 and 11 km s−1.
Fig. 7. Map of the velocity-integrated13CO(2–1) intensity in the HH 2
region. First contour and contour interval are 2 and 1 K km s−1 respec-tively. Velocity interval is indicated in the bottom right corner. The invidual HH knots are drawn in white.
We do not find any evidence of a large-scale velocity gra-dient in the cloud. The velocity field in the cloud and the “fila-ments” are rather smooth. The latter overlaps very well with the border of the cloud. On the contrary, there is a striking morphological association between the layers at blue velocities (5−8 km s−1) and the ambient cloud (see right panel in Fig. 1): the integrated emission of the fragment matches exactly the border of the ambient cloud, which is possible only because there is barely any ambient emission.
We detect two local maxima in the fragment downstream of HH 2: one coincides with the “molecular emission peak” and the other peaks at the center of the optical “ring” at offset position (15,−90) (see Fig. 8d). The gas in the fragment peaks at≈6.3 km s−1. The emission is characterized by bright lines (∼10 K) with a narrow linewidth (∼1 km s−1; see offset position (15, −75) in Fig. 2).
An LVG calculation performed with a temperature of 30 K gives a column density N(H2) 6.0 × 1021cm−2 at both
po-sitions in the layers. We note that the line opacity is quite high (0.68); it is probably the reason for the lack of contrast in the emission. The gas column density derived is in rough agree-ment with the determination obtained by Dent et al. (2003), who found N(H2)= 2 × 1022cm−2at the peak from dust
mil-limeter continuum observations. From the mean H2density in
the gas (∼2 × 105cm−3), we find that the thickness of the
emit-ting region is∼3−6×1016cm (5−10), much smaller than the
extent over the sky. The emission appears to be distributed in a sheet of dense gas rather than in a “round clump”. This means that the dense “ambient” gas emission is distribued around a cavity.
4.2. Shocked H2 around the cavity
Our ISOCAM/CVF observations (Fig. 8d) reveal some weak H2 emission in the S(2) and the S(3) lines above the 4σ
level in the gas shell where the second13CO maximum is
de-tected – offset position (15,−90). The H2 emission is
de-tected at the border of the gas layer, where the extinction is
detected. Therefore, we conclude that the S(2) and S(3) lines are tracing shocked gas over the shell. The non-detection of higher-J H2transitions suggests low-excitation conditions,
per-haps due to low-velocity shocks or previous (old) shocks. This strengthens our hypothesis that the shell of dense gas has been shaped by shocks, most likely from outflows.
It is therefore no wonder that the mid-IR emission map shows no correlation with the millimeter thermal dust emis-sion, which traces the “cold” dust. This component is detected mainly in the protostellar condensation VLA 1–4 and in the layers ahead of HH 2 (Dent et al. 2003). The ring and the jet co-incide with a minimum in the “cold” dust emission. Conversely, the maximum of absorption in the layer ahead of HH 2 pre-vents any mid-IR radiation from escaping the border of the cavity.
4.3. Dust grain properties
The dust grains in the ring exhibit properties very different from those observed in the protostellar envelopes of VLA 1–4. In the ring, we do not find any evidence at all of the presence of man-tle ices. The silicate absorption is so large towards VLA 1 that it makes it difficult to analyse the mantle composition. The lower absorption in the envelope of VLA 4 allows us to characterize the composition of the mantle ices: H2O, CH3OH, CH4, CO2
(Cernicharo et al. 2000). A broad bump from 11 to about 14µm is detected all over the ring. Such a bump is interpreted as the signature of crystalline silicates, although the exact composi-tion would require complementary data at longer wavelengths. Note that crystalline silicates have also been detected around the CS star (Cernicharo et al. 2000). Standard dust grains are expected to undergo deep changes in their composition when crossing shock(s) in the protostellar jet, as they release their icy mantle into the gas phase, either because of sputtering or shattering (Jones et al. 1994). Such processes result in an en-richment of the population of very small grains. Most likely, the very small grains detected by ISOCAM in the jet and in the ring are the result of this enrichment. In this context, a simple explanation for the detection of crystalline silicates is that they were originally present in grain cores and had already started to crystallize before the mantles evaporated. Further spectroscopy with the instruments onboard SPITZER should allow us to bet-ter constrain the composition and the physical paramebet-ters of the dust grains. As discussed above, the peculiar gas composition ahead of HH 2 most likely results from shock(s) too. Molecular line observations at higher angular resolution should be under-taken to better constrain the parameters of the shock.
Fig. 8. a) Mid-infrared emission integrated between 5 and 8.5 µm superposed on an optical [S
2
] image of HH 2 (Reipurth 1993). Contours range from 3 to 5.6 mJy by step of 0.2 mJy/pixel (px), from 6 to 10 mJy by step of 1 mJy/px, from 15 to 20 mJy by step of 5 mJy/px. b) High-velocity COJ= 2 → 1 emission (white contours) superposed on the 5.0−8.5 µm integrated mid-infrared map (colourscale). Contours range from 4 to 20 K km s−1by step of 2 K km s−1. The position of the individual shocks A-L are marked with stars. Red polygons delineate the fields used to compute individual spectra. c) Spectra of the mid-infrared emission between 5.0 and 17µm. d) Emission detected towards the “ring” West of HH 2H in:13CO(2–1): contours range from 6 to 11 K km s−1by step of 1 K km s−1; H2S(2) 12.2 µm: contours range from 4 to 10 mJy
by step of 2 mJy/px; 13.9−15.4 µm continuum: first contour is 4 mJy, following contours range from 4.5 mJy to 7.5 mJy by step of 1 mJy/px; PAH 11.3 µm band: contours range from 2 to 4.5 mJy by step of 0.5 mJy/px. The CVF maps were convolved with a Gaussian of 12HPFW to outline the extended emission. The rms in the 11.3 µm PAH map and the continuum map is 0.3 mJy/px; it is ≈1 mJy/px in the H2S(2) map.
of the sky. Both gas components (the fragment and the outflow wing) are blueshifted with respect to the ambient cloud.
On the other hand, the momentum carried away by the blueshifted outflow wind is≈0.1 Mkm s−1, i.e. a factor of a few less than the momentum necessary to detach the fragment from the cloud. The sum of several ejections in the past could have provided enough momentum; we speculate that another outflow, much more powerful, could be responsible for digging the cavity in the ambient cloud. When looking at the velocity channel map of SO emission (Fig. 3), it is interesting to note that the filament southwest of the core: a) consists of dense gas (SO); b) is velocity-shifted with respect to the ambient cloud; c) closely follows the border of the latter. Such a configura-tion is suggestive of a shock front propagating from the border of the cavity into the cloud. The orientation of such an out-flow would probably be closer to North-South than the actual HH 1–2 jet.
5. Photon-Dominated Region in HH 2
The emission integrated between 5 and 17 µm is shown in Figs. 8a, b. The SNR rms is∼0.1 mJy in the map. Protostar VLA 4 (Cernicharo et al. 2000) is detected as a bright point source in Figs. 8a, b. After subtracting the contribution of the ambient cloud, estimated from a reference position, the resid-ual flux appears very weak; a spectrum of the emission aver-aged over a reference field centered at (+100, −35) shows a contribution less than∼1 mJy (panel A in Fig. 8c). Hence, the emission discussed here arises from the HH 2 region. Southeast of the protostellar core down to HH 2, we detect mid-infrared emission along the HH 2 jet. The average flux level was es-timated over field D and is about 5 mJy longwards of 10µm (5σ above the rms noise level). The spatial coincidence with the CO high-velocity outflow (Fig. 1) suggests a physical asso-cation with the latter. The nature of the CVF emission appears to be mainly continuum in the jet, as indicated by the spectrum of Field D (Fig. 8c). This is the first time that mid-infrared
con-tinuum emission is reported towards a protostellar jet.
The map of the pure continuum between 13.9 and 15.4 µm shows that most of the flux comes from a ring-like structure (hereafter the “ring”) of 60diameter (0.13 pc), westwards of HH 2H (Fig. 8d), This annular structure is detected in the op-tical SII line (Reipurth et al. 1993) and in the 5.0−8.5 µm emission map. The minimum of emission in the center of the ring coincides with the peak of gas column density detected in13CO, at offset position (15,−75). This provides direct
ev-idence that the mid-IR “ring” and the molecular gas surround-ing the cavity are closely related. The average spectrum of the
is no hint of PAH emission bands. The relatively low SNR and the strong cloud extinction prevent us from drawing a conclu-sion about the presence/absence of PAHs along the jet.
A rough estimate of the FUV field intensity is obtained by integrating the flux in the ISOCAM band. This yields an infrared luminosity L ∼ 0.01 L, and a FUV field:
G= 20 G0. This is fully consistent with the estimate obtained
by Molinari & Noriega-Crespo (2002) from the analysis of the FIR [CII] 158µm and [OI] 63µm, 145 µm lines and the FIR continuum observed at 80resolution with ISO/LWS. We note that our direct estimate of the FUV lies in the low range of val-ues (20–1000 G0) required by the models of UV-driven
photo-chemistry (Viti et al. 2003) to account for the molecular emis-sion ahead of HH 2.
Continuum emission in the jet and the ring could be fitted by a greybody with an opacity lawτν ∝ ν and a temperature
Td= 200−250 K. Assuming that the emission comes from
typ-ical interstellar a(big) dust grains, we find that the amount of material corresponds to an Av 0.1−1 × 10−6, which
trans-lates into very low gas column densities. The efficiency of dust grains with a size as small as 0.003 µm is too low (Draine & Lee 1984) to reach equilibrium temperatures of this order, un-less the UV field is actually very strong, far above our esti-mate. Most likely, we are observing the emission of very small dust grains, transiently heated to high temperatures by the local UV field.
Weak continuum flux is also detected across the HH ob-ject and West of the protostellar core VLA 1–4. This extension coincides with the border of the cavity in the molecular cloud.
5.1. HH 2H as exciting source of the PDR
Previous observations in the UV have revealed strong ex-tended emission in the direction of HH 2H-2A (Böhm-Vitense et al. 1982; Raymond et al. 1997). The other sources in the field detected by IUE are HH 1 and the Cohen-Schwartz star; both are located too far away to account for the PDR ob-served around HH 2. Analysis of optical high-ionization lines has shown evidence for a J-type shock with velocities of ∼180 km s−1 moving into previously accelerated gas
(opti-cal measurements indicate proper motions up to 400 km s−1 towards HH 2H and HH 2A (Bally et al. 2002), identifying HH 2H as the actual jet impact region (or Mach disk region). HH 2H-2A is therefore the most plausible candidate as an ex-citing source; it is also the only place where X-ray emission has been detected in the region (Pravdo 2001).
HUT measurements by Raymond et al. (1997) show that the UV continuum is dominated by 2 photon radiation of col-lisionally excited H atoms longwards of 1750 Å. At shorter wavelengths, the continuum is dominated by H2 Lyman
emis-sion bands. Once corrected for the interstellar extinction, the total UV luminosity is∼0.3 L; it is distributed roughly equally between both mechanisms. The exact solid angle encompassed by the ring is difficult to estimate. Following our hypothesis that the outflow has dug a cavity around HH 2, it is reason-able to assume that the ring and the molecular emission peak ahead of HH 2 are at about the same distance to the UV source, knots H-A (50). The diameter of the ring is 60, hence the fraction of photons intercepted by the ring is∼5%, which com-pares well with the mid-IR luminosity radiated by the very small grains in the walls of the cavity (3% of the UV lumi-nosity produced in HH2H-2A).
The possibility of a similar mechanism (strong J-shocks) to account for the mid-infrared emission along the protostellar jet, between VLA 1 and HH 2, should be explored. Deep sensitive spectroscopy of the entrained gas could allow one to constrain such possibility. Another possibility is that the jet region is ac-tually almost free of dust grains after the crossing of the strong shocks which are now impacting HH 2. In this case we would be detecting the walls of a dust-free cavity around the jet, illu-minated either by HH 2H or by the powering source.
6. Summary and conclusions
The observations presented in this paper allow us to draw a more complete and complex picture of the HH 2 region. In addition to the high-velocity molecular outflow associated with the optical jet HH 1–2, we have detected another out-flow component that propagates over 50in the molecular gas downstream HH 2. This outflow is associated with a weak jet detected in the optical, which appears to come from HH 2. The orientation of this component in the sky differs from the HH 2 jet and could be the result of jet deflection on the dense obstacles responsible for the knots of H2 emission in HH 2.
Another possibility is that this component traces a previous ejection from the source, which would be precessing. High an-gular resolution observations should allow us to test the first hypothesis, by comparison with numerical modelling (see e.g. Raga & Canto 1995)
This provides direct evidence that the molecular gas ahead of HH 2 has been affected by protostellar ejections. It is con-sistent with the detection of optical bow-shocks at large dis-tances from the HH 1–2 region. Several kinematical compo-nents are detected in the gas downstream of HH 2. Observations of SO sreveal the presence of very dense gas (∼2−5×105cm−3)
associated with the low-velocity shocks along the HH 2 jet, and in the blueshifted outflow, close to knot HH 2L. The latter po-sition coincides with the peak of emission in the region. We conclude that SO is probably tracing the shock interaction of the outflow with the ambient gas.
The gas layers downstream of HH 2 exhibit evidence of shocks: East of HH 2, ISOCAM observations detect H2 line
emission in the gas layers, which is the signature of shocks, probably old enough to be detected only in the low-excitation
transitions S(2) and S(3). The gas layers have been shaped in a shell, probably as the result of prostellar outflow interaction with the cloud. Actually, the morphology and kinematics of the filament Southwest of the protostellar core are compatible with their tracing a shock-compressed“back side” of the cavity (the “front side” being the gas shell downstream of HH 2).
ISOCAM observations show emission from very small grains at the inner surface of the cavity downstream of HH 2. This could result from the shattering and sputtering of large in-terstellar dust grains in the outflow/cloud interaction. We find that the UV field produced in the strong shock HH 2H-A, which illuminates the inner side of the cavity, creates a Photon-Dominated Region of FUV intensity G∼ 20 G0. It is sufficient
to account for the emission of the very small grains and the PAHs detected. Mid-infrared emission along the jet remains a puzzle; deep mid-infrared spectroscopy along the jet could al-low one to determine the origin of the emission, in particular if it is related to strong dissociative shocks between the source and HH 2.
Our observations indicate that some of the molecular species whose abundance is found enhanced in the gas down-stream of HH 2 are actually produced in the region of outflow interaction with the cloud. Mapping at better angular resolu-tion, comparable to the data presented here, should be under-taken. In combination with detailed shock modelling, it would allow us to explore the shock hypothesis and estimate its rel-ative contribution with respect to UV-induced photodesorption in the gas phase enrichment of “unusual” molecular species.
Acknowledgements. J. Cernicharo acknowledges the Spanish DGES
for this research under grants AYA2000-1784 and AYA2003-2784. We thank Dr. F. Boulanger for many stimulating discussions on the ISOCAM observations of HH 2.
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![Fig. 8. a) Mid-infrared emission integrated between 5 and 8.5 µm superposed on an optical [S 2 ] image of HH 2 (Reipurth 1993)](https://thumb-eu.123doks.com/thumbv2/123doknet/14802131.606727/9.892.128.821.105.967/fig-infrared-emission-integrated-superposed-optical-image-reipurth.webp)
